MULTI-WAVELENGTH ANALYSIS OF THE GALACTIC
SUPERNOVA REMNANT MSH 11-61A
The MIT Faculty has made this article openly available. Please share
how this access benefits you. Your story matters.
Citation
Auchettl, Katie, Patrick Slane, Daniel Castro, Adam R. Foster,
and Randall K. Smith. “MULTI-WAVELENGTH ANALYSIS OF
THE GALACTIC SUPERNOVA REMNANT MSH 11-61A.” The
Astrophysical Journal 810, no. 1 (August 27, 2015): 43. © 2015
The American Astronomical Society
As Published
http://dx.doi.org/10.1088/0004-637x/810/1/43
Publisher
IOP Publishing
Version
Final published version
Accessed
Wed May 25 22:57:25 EDT 2016
Citable Link
http://hdl.handle.net/1721.1/99908
Terms of Use
Article is made available in accordance with the publisher's policy
and may be subject to US copyright law. Please refer to the
publisher's site for terms of use.
Detailed Terms
The Astrophysical Journal, 810:43 (12pp), 2015 September 1
doi:10.1088/0004-637X/810/1/43
© 2015. The American Astronomical Society. All rights reserved.
MULTI-WAVELENGTH ANALYSIS OF THE GALACTIC SUPERNOVA REMNANT MSH 11-61A
Katie Auchettl1,2, Patrick Slane1, Daniel Castro3, Adam R. Foster1, and Randall K. Smith1
1
Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138, USA
School of Physics and Astronomy, Monash University, Melbourne, Victoria 3800, Australia
MIT-Kavli Center for Astrophysics and Space Research, 77 Massachusetts Avenue, Cambridge, MA 02139, USA
Received 2015 May 1; accepted 2015 July 28; published 2015 August 27
2
3
ABSTRACT
Due to its centrally bright X-ray morphology and limb brightened radio profile, MSH 11–61A (G290.1–0.8) is
classified as a mixed morphology supernova remnant (SNR). H I and CO observations determined that the SNR is
interacting with molecular clouds found toward the north and southwest regions of the remnant. In this paper we
report on the detection of γ-ray emission coincident with MSH 11–61A, using 70 months of data from the Large
Area Telescope on board the Fermi Gamma-ray Space Telescope. To investigate the origin of this emission, we
perform broadband modeling of its non-thermal emission considering both leptonic and hadronic cases and
concluding that the γ-ray emission is most likely hadronic in nature. Additionally we present our analysis of a
111 ks archival Suzaku observation of this remnant. Our investigation shows that the X-ray emission from MSH
11–61A arises from shock-heated ejecta with the bulk of the X-ray emission arising from a recombining plasma,
while the emission toward the east arises from an ionizing plasma.
Key words: cosmic rays – gamma rays: ISM – ISM: individual objects (G290.1-0.8) – ISM: supernova remnants –
X-rays: ISM
limb-brightened (Rho & Petre 1998). The evolutionary
sequence leading to these unusual X-ray properties are not
well understood and the morphology and characteristics of
these SNRs are difficult to explain using standard SNR
evolution models. There are two main models that are invoked
in the literature to explain their characteristics. One possible
model (White & Long 1991) assumes that in the vicinity of the
supernova explosion there are many small, dense, cold
cloudlets. These cloudlets are small enough that they do not
affect the passage of the shock, and are sufficiently dense that
they are neither blown apart nor swept up. Once the shock has
passed, the cloudlets slowly evaporate, filling the interior of the
SNR with a relatively dense gas that emits in X-rays. Another
possible scenario is that thermal conduction results in the
transport of heat and material to the center of the remnant,
increasing the central density of the remnant, and smoothing
the temperature gradient behind the shock (Cox et al. 1999).
The ionization state of a thermal plasma in an SNR can be
characterized by its ionization temperature (TZ) which
describes the extent to which the ions are stripped of their
electrons and its electron temperature (Te) which describes the
kinetic energy of the electrons. The thermal plasmas of SNRs
have been thought to be either underionized, where TZ < Te or
in collisional ionization equilibrium TZ = Te. Recent observations by the Suzaku satellite have confirmed earlier suggestions
based on Advanced Satellite for Cosmology and Astrophysics
(ASCA) data (Kawasaki et al. 2005), that the thermal plasma in
some MM SNRs is overionized (recombining; e.g., 3C391:
Sato et al. 2014; Ergin et al. 2014). Recombining plasmas have
ionization temperatures that are higher than the electron
temperatures and require rapid cooling of electrons either by
thermal conduction (Kawasaki et al. 2002), adiabatic expansion
via rarefraction and recombination (Itoh & Masai 1989) or the
interaction with dense cavity walls or MCs (Dwarkadas 2005).
MSH 11–61A (G290.1–0.8) is a Galactic MM SNR that is
known to be interacting with MCs. It was first discovered by
Mills et al. (1961) using the Sydney 3.5 m cross-type radio
1. INTRODUCTION
Since the launch of the Large Area Telescope (LAT)
onboard the Fermi Gamma-ray Space Telescope, its improved
sensitivity and resolution in the MeV–GeV energy range has
lead to a number of supernova remnants (SNRs) being detected
in GeV γ-rays. The shock-front of an SNR is expected to be
able to accelerate cosmic rays (CRs) efficiently, producing nonthermal X-ray and γ-ray emission. As γ-rays can arise from
leptonic processes such as inverse Compton (IC) scattering and
non-thermal bremsstrahlung from high-energy electrons, or
from hadronic emission arising from the decay of a neutral pion
(produced in a proton–proton interaction) into two photons, a
means of distinguishing between these two mechanisms is
crucial for our understanding of the origin of this observed
emission. Thermal and non-thermal emission from SNRs have
provided increasing support in favor of CRs being accelerated
at the shock front of the remnant (e.g., Tycho: Warren et al.
2005; RX J1713.7-3946: Uchiyama et al. 2007; W44, MSH 1739 and G337.7-0.1: Castro et al. 2013). SNRs known to be
interacting with dense molecular clouds (MCs) are ideal,
indirect laboratories that one can use to detect and analyze γrays arising from accelerated protons. The interaction of the
SNR’s shockwave with dense molecular material is often
inferred from the detection of one or more OH (1720 MHz)
masers, but enhancement of excitation line ratios such as J = 2
1 J = 1  0 , broadenings and asymmetries in molecular
line features or morphological alignment of molecular features
with SNR features can also allow one to determine SNR/MC
interaction (see Slane et al. 2015 and references therein).
Some of the first SNRs detected by the Fermi-LAT (e.g.,
W44: Abdo et al. 2010b; Ackermann et al. 2013; IC443:
(Ackermann et al. 2013); 3C391: Castro & Slane 2010; and
W49B: Abdo et al. 2010a) are part of a unique class called
Mixed-Morphology (MM) SNRs. Some of these SNRs are
known to be interacting with MCs. These SNRs are
characterized by their centrally peaked X-ray morphology
which is thermal in nature, while their radio profiles are
1
The Astrophysical Journal, 810:43 (12pp), 2015 September 1
Auchettl et al.
telescope. It was first identified as an SNR by Kesteven (1968)
and later classified as a shell-type SNR with a complex internal
structure and ear-like protrusions toward the northwest and
southeast using the Molongo Observatory Synthesis Telescope
(MOST) at multiple different wavelengths (Kesteven &
Caswell 1987; Milne et al. 1989; Whiteoak & Green 1996).
It has an angular size of 19′ × 11′ and a radio-continuum
spectral index of α = −0.33+/−0.07 (Reynoso et al. 2006).
Radio continuum observations using the Australia Telescope
Compact Array by Reynoso et al. (2006) showed filamentary
emission with little shell structure, while the northern and
southern edges of the remnant show evidence that the shock
front could be interacting with a plane parallel density gradient.
Using NANTEN CO images of MSH 11–61A, Filipovic et al.
(2005) determine that the SNR is associated with a MC toward
the south–west and northern rim of the remnant. H I
observations using the Southern Galactic Plane Survey find
that the MC is found at a local standard of rest velocity of
∼13 km s−1 (McClure-Griffiths et al. 2005).
MSH 11–61A was first detected in X-rays by Seward (1990)
using a 10.9 ks Einstein Observatory observation. The
0.3–4.5 keV Imaging Proportional Counter (IPC) image of
the remnant showed that the X-ray emission is peaked toward
the center of the remnant. Using a 40 ks ASCA GIS
observation, Slane et al. (2002) were able to determine that
the central X-ray emission is thermal in nature, classifying the
remnant as a MM SNR. They modeled the X-ray emission
using the cloudy ISM model by White & Long (1991) and
derived an intercloud medium density of ∼0.05–0.40 cm−3 and
an age of 10–20 kyr. Using XMM-Newton, García et al. (2012)
analyzed five regions along the axes of the remnant using an
absorbed plane parallel non-equilibrium ionization (NEI)
(VPSHOCK) model and found that the physical conditions
across the remnant are not homogeneous, with variation in
ionization state, temperature and elemental abundances.
Kamitsukasa et al. (2015) analyzed Suzaku data and found
that in the center and in the northwest of the remnant the
plasma is recombining, while everywhere else it is ionizing.
The distance to MSH 11–61A has been measured by a
number of different methods. H I measurements taken by the
64 m Parkes telescope (Goss et al. 1972; Dickel 1973) gives a
lower limit of ∼3.5 kpc to the remnant. Hα measurements by
Rosado et al. (1996) using a Fabry–Perot interferometer
implied a distance of 6.9 kpc assuming a VLSR of +12 km s−1
for the SNR. Using CO measurements, Reynoso et al. (2006)
derived a distance of 7–8 kpc assuming the Brand & Blitz
(1993) rotation curve. Reynoso et al. (2006) derived a distance
of 7 ± 1 kpc using H I absorption measurements from ATCA,
combined with data from the Southern Galactic Plane Survey,
while Slane et al. (2002) estimated a distance of 8–11 kpc by
modeling the thermal X-ray emission of the remnant as
detected by ASCA. We use 7 kpc throughout this paper.
There are three pulsars close to the position of MSH
11–61A. Kaspi et al. (1997) discovered the young (spin-down
age, τ = 63 kyr), energetic pulsar PSR J1105-6107 (J1105)
which is located approximately 25′ away from the remnant. It
has a spin-down luminosity of 2.5 × 1036 erg s−1 and overlaps
the position of the EGRET γ-ray source 3EG J1103-6106. It
was also detected via periodicity searches in GeV γ-rays by the
Fermi-LAT satellite (Abdo et al. 2013). It has a dispersion
measure of 271 cm−3 pc which implies a distance of ∼7 kpc
using the standard Galactic electron density model (Cordes &
Lazio 2002). Kaspi et al. (1997) considered the scenario that
this pulsar is associated with the remnant and determined from
proper motion measurements that it would need to be travelling
with a transverse velocity of ∼650 km s−1 to have reached its
current position, assuming a distance of 7 kpc to the pulsar and
τ = 63 kyr. This is much larger than the average pulsar
transverse velocity but much less than what has been suggested
for other pulsar-SNR associations (e.g., Caraveo 1993), leading
the authors to conclude that association is possible. Using the
ASCA X-ray characteristics of MSH 11–61A, Slane et al.
(2002) concluded that MSH 11–61A and J1105 are not
associated under the assumption that the SNR evolved via
thermal conduction or a cloudy ISM. These two models imply
a transverse velocity of ∼4.5 × 103 km s−1 and ∼5.3 ×
103 km s−1 respectively, which is much larger than the mean
velocity (∼310 km s−1) of young pulsars (Hobbs et al. 2005).
The two other pulsars, PSR J1103-6025 (Kramer et al. 2003)
and PSR J1104-6103 (Kaspi et al. 1996) are not associated with
the remnant as they have characteristic ages much larger than
1 Myr, which is far greater than the expected lifetime of an
SNR. Also nearby is the extended INTEGRAL source IGR
J11014-6103, whose neutron star PSR J1101-6101 (Halpern
et al. 2014) is traveling at a velocity exceeding 1000 km s−1,
which Pavan et al. (2014) associate with MSH 11–61A.
Using 70 months of Fermi-LAT data, we analyze the GeV γray emission coincident with MSH 11–61A and investigate the
nature of this emission using broadband modeling. In addition,
we analyze archival Suzaku data and report on the spatial and
spectral properties of the X-ray emission of this remnant. In
Section 2 we describe how the Fermi-LAT data are analyzed
and present the results of this analysis. In Section 3 we present
our spatial and spectral analysis of the Suzaku observation of
MSH 11–61A, while in Sections 4 and 5 we discuss the
implications of our results.
2. OBSERVATIONS AND ANALYSIS OF MSH 11–61A
We analyzed ∼70 months of reprocessed data collected by
the Fermi-LAT from 2008 August 4 to 2014 June 16. We
selected data within a radius of 20° centered on MSH 11–61A.
We used the “P7REP_SOURCE_V15” instrument response
function (IRF) which is based on the same in-flight event
analysis and selection criteria that was used to generate the
previous “PASS7_V6” IRFs (details described in Ackermann
et al. 2012). Due to the improved reconstruction of the
calorimeter position as well as a 1% per year correction for the
degradation of the light yield of the calorimeter, the new IRFs
significantly improve the point-spread function (PSF) of the
LAT for energy >5 GeV.4 The systematic uncertainties in the
effective area of the Fermi-LAT using “P7REP_SOURCE_V15” are 10% below 100 MeV, decreasing logarithmically
in energy to 5% between 0.316 and 10 GeV and increasing
logarithmically to 15% at 1 TeV.5 We selected events with a
zenith angle less than 100° and that were detected when the
rocking angle of the LAT was greater than 52° to decrease the
effects of terrestrial albedo γ-rays. We analyzed the γ-ray data
in the direction of MSH 11–61A using the Fermi Science Tools
v9r33p0.6
4
More details found: http://www.slac.stanford.edu/exp/glast/groups/
canda/lat_Performance.htm
5
http://fermi.gsfc.nasa.gov/ssc/data/analysis/LAT_caveats.html
6
http://fermi.gsfc.nasa.gov/ssc/data/analysis/software/
2
The Astrophysical Journal, 810:43 (12pp), 2015 September 1
Auchettl et al.
Due to the low count rates of γ-rays and the large PSF of the
Fermi-LAT, the standard maximum likelihood fitting technique, gtlike, was used to analyze the γ-ray emission of the
remnant. Given a specific emission model, gtlike determines
the best-fit parameters of this model by maximizing the joint
probability of obtaining the observed data given an input
model. Gtlike accounts for background γ-ray emission by using
diffuse Galactic and isotropic emission models described by the
mapcube file gll_iem_v05_rev1.fits and the isotropic
spectral template iso_source_v05_rev1.txt.7 Gammaray emission from sources found in the Fermi-LAT second
source catalog are fixed to their position listed in the catalog
and their background contribution is calculated. The Galactic
diffuse emission arises from the interaction of CRs with the
interstellar medium and their subsequent decay into γ-rays,
while the isotropic component arises from diffuse extragalactic
γ-rays and residual charged particle emission.
To improve the angular resolution of the data while
analysing the spatial properties of the γ-ray emission of MSH
11–61A, we selected γ-ray data converted in the front section
of the Fermi-LAT with an energy range of 2–200 GeV. The
improvement in spatial resolution in this energy range arises
from the fact that the 1σ containment radius angle for frontselected photon events is <0 °. 3, while for lower energies it is
much larger. To determine the detection significance, position
and possible extent of the γ-ray emission coincident with
MSH 11–61A we produced test statistic (TS) maps using
gttsmap with an image resolution of 0 °. 05. The TS is defined
as 2 log (L ps L null ), where Lps is the likelihood of a point
source being found at a given position on a spatial grid and
Lnull is the likelihood of the model without the additional
source.
To determine the spectral energy distribution (SED) of the γray emission coincident with MSH 11–61A we use events
converted in the front section of the LAT that have an energy of
0.2–204.8 GeV. This energy range is chosen to avoid the large
uncertainties in the Galactic background model that arise below
0.2 GeV and to reduce the influence of the rapidly changing
effective area of the LAT at low energies. We model the flux in
each of 8 logarithmically spaced energy bins and estimate the
best-fit parameters of the data using gtlike. We also include in
the likelihood fit, background sources from the 24 month
Fermi-LAT second source catalog (Nolan et al. 2012) that are
found within 20° region centered on MSH 11–61A. All evident
background sources were identified in the Fermi-LAT second
source catalog and the associated parameters from the catalog
were used. We left the normalization of the Galactic diffuse
emission, isotropic component and the background point
sources within 5° of MSH 11–61A free. For all other
background point sources with a distance greater than 5° from
MSH 11–61A their normalizations were frozen to that listed in
the 2nd Fermi-LAT catalog. In addition to the statistical
uncertainties that were obtained from the likelihood analysis,
systematic uncertainties associated with the Galactic diffuse
emission were also calculated by artificially altering the
normalization of this background by ±6% from the best-fit
value at each energy bin as outlined in Castro & Slane (2010).
In Figure 1 left panel, we have generated a 2 °. 5 × 1 °. 5 count
map centered on MSH 11–61A that is smoothed by a Gaussian
with a width similar to the PSF for the events selected. MSH
11–61A is located in a very complicated region of the sky.
There are many Fermi-LAT sources close to the remnant, with
γ-ray bright SNR MSH 11-62 (2FGL J1112.1-6040 and 2FGL
J1112.5-6105) located ∼1 °. 2, and PSR J1105-6107 (2FGL
J1105.6-6114) located ∼0 °. 36, from the remnant. There are
four other Fermi-LAT sources in the immediate vicinity of
MSH 11–61A (2FGL J1104.7-6036, 2FGL J1105.6-6114,
2FGL J1059.3-6118c and 2FGL J1056.2-6021), but none of
these are coincident with the MOST radio contours of the SNR.
We define a source region at the position of MSH 11–61A to
estimate the flux from the SNR. Due to the close proximity of
PSR J1105-6107 (J1105) and the low resolution of the FermiLAT PSF (∼1° for front events at 68% containment at
∼0.6 GeV), we cannot rule out that this pulsar is not
contributing significantly to the observed γ-ray emission
seen in Figure 1. Thus to analyze the spatial and spectral
characteristics of the γ-ray emission of MSH 11–61A, we have
to remove the pulsar contribution. As MSH 11-62 is located
>1° away from MSH 11–61A, it is unlikely that it is
contributing significantly to the observed γ-ray emission.
2.1. Removing the γ-Ray Contribution of PSR J1105-6107
To avoid contamination from the pulsed emission of PSR
J1105-6107 (J1105), we must perform our analysis in the offpulse window of the pulsar light curve. The Fermi-LAT
collaboration made available with its second Fermi-LAT
catalog of γ-ray pulsars (Yu et al. 2013) the ephemerides of
117 γ-ray emitting pulsars that had been detected using three
years of Fermi-LAT data. Due to the continuous observations
of the Fermi-LAT, Yu et al. (2013) were able to directly
determine regular times of arrival (TOAs) to produce a precise
pulsar ephemeris. We use the available ephemeris for J1105
which is valid from MJD 54200 (2007 April) to MJD 56397
(2013 April). As we are interested in analysing over ∼70
months of Fermi-LAT, we need to check the validity of using
the available J1105 ephemeris over the whole ∼70 months.
Using the Parkes telescope, Yu et al. (2013) analyzed J1105
searching for timing irregularities over a period of ∼16 years
from 1994 August to 2010 September (MJD 49589–MJD
55461). J1105 experienced a glitch at MJD 50417 (1996
November), MJD 51598 (2000 February), MJD 54711 (2008
September) and MJD 55288 (2010 April). The Fermi-LAT
ephemeris of J1105 covers the glitches experienced by the
pulsar since the beginning of the Fermi-LAT mission. As of
writing, there have been no other reports in the literature that
J1105 has undergone glitches since 2010 April. We also tested
whether the pulsar is noisy in γ-rays, as this can also indicate
irregularities in the pulsar rotation that have not been
incorporated in the ephemeris. We produced a pulse phase
versus events plot during the period of validity of the
ephemeris (MJD 54200 to MJD 56397) and for the period
between the end time of the ephemeris and the end date of
our data (MJD 56397 to MJD 56824) to see if the pulse peaks
that we observe in Figure 2 disappear due to noise of the
pulsar. We find that during both time periods we obtain the
same pulse phase profile. As no glitches have been detected
as of writing and the pulsar has not been noisy beyond
MJD 56397, we assume that this ephemeris is valid for our
whole data set.
The γ-ray photons were phase-folded using this ephemeris
and the pulse phases were assigned to the Fermi-LAT data
using the Fermi-LAT TEMPO2 plugin provided by the Fermi-
7
The most up to date Galactic and isotropic emission models can be found
http://fermi.gsfc.nasa.gov/ssc/data/access/lat/BackgroundModels.html
3
The Astrophysical Journal, 810:43 (12pp), 2015 September 1
0.053
0.091
0.13
Auchettl et al.
0.17
5.4
11
16
22
27
−2
Figure 1. Left: 2 °. 5 × 1 °. 5 Fermi-LAT count map of the 2–200 GeV γ-ray emission surrounding MSH 11–61A (units: counts degree ). The pixel binning is 0 °. 01
and the maps are smoothed with a Gaussian of width 0 °. 2. The cyan circles correspond to the background Fermi-LAT sources surrounding MSH 11–61A. The
brightest emission seen to the east of MSH 11–61A corresponds to the γ-ray bright SNR MSH 11-62. The yellow X corresponds to the radio position of PSR J11056107, while the white X corresponds to the position of ICG J11014-6103 (the lighthouse nebula). One can see that MSH 11–61A is located in a complicated region of
the γ-ray emitting sky. 843 MHz MOST radio continuum contours are overlaid in green. Right: 2 °. 5 × 1 °. 5 TS map of MSH 11–61A. The magenta contours
corresponds to the H I emission of the MC associated with the remnant as detected by the Southern Galactic Plane Survey (McClure-Griffiths et al. 2005). The H I
emission contours range from 66 Kelvin at the edge of the molecular cloud interacting with MSH 11–61A to 100 Kelvin at the center of the molecular cloud over
seven linearly spaced contour levels.
components, we include in the background model the FermiLAT sources associated with MSH 11-62 (2FGL J1112.1-6040
and 2FGL J1112.5-6105) and the four sources in the immediate
vicinity of the remnant (2FGL J1104.7-6036, 2FGL J1105.66114, 2FGL J1059.3-6118c and J1056.2-6021). The TS map
suggests that there is significant γ-ray emission coincident
with MSH 11–61A and the MC associated with the remnant,
as highlighted by the magenta contours. The peak of
the γ-ray emission is found at a best fit position of
(aJ2000 , dJ2000 ) = (11h 01m 29s , -605529), placing it outside
the SNR boundary, but consistent with being located along or
inside the western limb given the angular resolution of the
Fermi-LAT. The emission is detected with a significance of
∼5σ, with the γ-ray emission at the center of the remnant
producing a significance of ∼4σ. One can see in Figure 1 right
panel that the contribution of MSH 11-62 and the other FermiLAT sources surrounding the remnant have been modeled out,
while the contribution from the pulsar (J1105) has been gated
out successfully.
Figure 2. Pulse-phase diagram of PSR J1105-6107 obtained using events in an
energy range of 0.1–300 GeV coming from a 0 °. 5 radius around the position of
the pulsar. The off-pulse window used for this analysis is defined between
0.00–0.05, 0.22–0.55 and 0.70–1.00 of the pulse phase. Here two cycles are
shown.
2.3. γ-Ray Spectrum
LAT collaboration.8 This plugin calculates the rotational phase
of the pulsar for each photon arrival time in the Fermi-LAT
data using the barycentric dates of each event. Using ftselect,
we remove the pulse and use only the γ-ray photons in the offpulse window (defined by the 0.00–0.05, 0.22–0.55 and
0.70–1.00 pulse phase intervals) to perform our spectral and
morphological analysis. In Figure 2 we have plotted the pulsephase diagram of J1105 obtained in the 0.10–300 GeV energy
range using 0 °. 5 radius around the position of J1105.
The γ-ray spectrum of MSH 11–61A is shown in Figure 3,
with the statistical errors plotted in black and systematic errors
plotted in red. For energies above 6.40 GeV, only flux upper
limits have been determined and are plotted as blue triangles.
Additionally, the best fit power law and exponential cut-off
power law models are plotted as the green dotted line and
purple dotted–dashed line respectively. The γ-ray spectrum can
be fit using a simple power law with a spectral index of
2
+0.07
2.750.06 , giving a reduced χ ∼ 1. An exponential cut-off
+1.91
power law with Ecut = 4.20-0.66 GeV and spectral index of
2
+0.17
2.490.19 can fit the spectrum equally well giving a reduced χ
∼ 0.8 for the fit. For an energy range of 0.1–100 GeV, the
+0.34
-11 erg cm−2 s−1, assuming
integrated flux is (4.20.98 ) ´ 10
the power law fit. Using a distance of 7 kpc, the luminosity of
+0.17
35
this γ-ray source in this energy range is (2.50.63 ) ´ 10
erg s−1.
2.2. TS Map
In Figure 1 right panel we present a TS map of MSH
11–61A using the off-pulse γ-ray data. This was calculated
using gttsmap over an energy of 0.2–2.0 GeV and using front
events only. In addition to the diffuse Galactic background
8
http://fermi.gsfc.nasa.gov/ssc/data/analysis/user/Fermi_plug_doc.pdf
4
The Astrophysical Journal, 810:43 (12pp), 2015 September 1
Auchettl et al.
spectral fitting package (XSPEC) version 12.8.2q with
AtomDB 3.0.112 (Smith et al. 2001; Foster et al. 2012).
3.1. Spectral Analysis of the Individual Annulus
and Rectangular Regions
We extracted spectra from a central circular region defined
by region 1 in Figure 4 (right panel) and 6 annular regions of
width 0′. 82 to cover the central X-ray emission of the remnant
(regions 2–7 in Figure 4, right panel). The radial size of these
regions was chosen to be the same size as the angular
resolution of Suzaku (∼0′. 8). These regions were chosen to
fully enclose the bright central X-ray emission of the remnant,
which is quite symmetric in nature. We also extracted spectra
from three rectangular regions (regions 8–10 in Figure 4, right
panel) that are not covered by the annulus regions, to enclose
protrusions in the northwest, southeast, and east. Annular
regions were chosen in order to characterize radial variations in
the brightness, temperature, ionization state and elemental
abundances of the remnant, all of which are important for
understanding the nature of the MM. Our choice of regions
differs from those chosen by Kamitsukasa et al. (2015), who
also analyzed the Suzaku observation of MSH 11–61A. They
extracted spectra from five regions that do not enclose the full
X-ray emission from the remnant—a central region that
corresponds to our regions 1, 2, and 3; a northwest region
that encompasses our region 8 and a northwestern portion of
region 7; and NE, SE, and SW regions that cover sectors of our
annular regions and also encompass our region 9 in the east.
All spectra were grouped by 20 counts using the FTOOLS
command grppha and all available XIS detectors were used.
To estimate the background, we extract data from the full
field of view of the XIS of our observation, excluding the
calibration regions and the emission from the remnant. The
background spectrum consists of two major components, the
non X-ray background and the astrophysical background which
is made up of the cosmic X-ray background, the Galactic ridge
X-ray emission and the Galactic halo. We use xisnxbgen (Tawa
et al. 2008) to generate a model for the NXB which we then
subtract from our background spectrum. Similarly, we subtract
a model for the NXB from our spectra obtained from the
regions shown in Figure 4 right panel. Similar to Kamitsukasa
et al. (2015), we model this NXB subtracted spectrum. We fix
the cosmic X-ray background power law component to that of
Kushino et al. (2002), and use a single apec model with a
temperature and surface brightness (∼1.0 × 10−12 erg
cm−2 s−1 deg−2 for 0.5–2.0 keV) similar to that obtained by
Henley & Shelton (2013) to define our Galactic halo
component. We use a single low temperature apec model with
subsolar abundances frozen to that of Kaneda et al. (1997) to
define the Galactic ridge emission. We also use the Wilms et al.
abundance table (Wilms et al. 2000). Our best-fit parameters
for our background and their uncertainties are given in Table 1.
To model the X-ray emission from the remnant we used a
NEI collisional plasma model, VVRNEI, which is characterized by a final (kT) and initial electron temperature (kTinit),
elemental abundances and a single ionization timescale
(τ = net). This allows one to model a plasma that is ionizing
up to collisional equilibrium from a very low initial
temperature kTinit, mimicking the standard NEI/VNEI model
that is commonly used in the literature. Additionally, the
RNEI/VVRNEI model can reproduce a recombining (overionized) plasma where one assumes that the plasma starts in
Figure 3. Fermi-LAT γ-ray spectrum of MSH 11–61A. Statistical errors are
shown as black error bars, systematic errors are plotted as red error bars, while
the upper limits are plotted as blue triangles. The simple power law model with
+0.17
+0.07
G = 2.750.06 and exponential cut-off power law model with G = 2.49-0.19
+1.91
and Ecut = 4.20GeV
are
shown
as
the
green
dotted
and
magenta
dotted–
0.66
dashed line respectively.
3. Suzaku OBSERVATIONS OF MSH 11–61A
MSH 11–61A was observed with Suzaku using X-ray
imaging spectrometers (XIS) (Koyama et al. 2007) on 2011
July 25th for ∼111 ks (ObsID 506061010). For this observation only XIS0,9 XIS1 and XIS3 observations are available as
XIS2 has not been functional since 2006 November.10
Recently, Kamitsukasa et al. (2015) presented their analysis
of this Suzaku observation. They found recombining plasma in
the center and in the northwest of the remnant which has
enhanced abundances and a temperature of ∼0.5 keV, while
everywhere else the X-ray emission arises from an ionizing
ISM component with a temperature of ∼0.6 keV. In Section 5.2.1, we estimate the density of the γ-ray emitting
material based on our Fermi-LAT spectrum (Figure 3). To test
whether this inferred density agrees with other observations, we
have re-analyzed the Suzaku data in order estimate the density
of the surrounding environment.
For our analysis we used the standard tools of HEASOFT
version 6.16. We reprocessed the unfiltered public data using
aepipeline (version 1.1.0) and use the current calibration
database (CALDB) available as of 2014 July 1st (version
20140701). Following the standard screening criteria,11 we
filtered hot and flickering pixels, time intervals corresponding
to Suzaku passing the South Atlantic Anomaly and night-Earth
and day-Earth elevation angles less than 5° and 20°,
respectively. We utilized events that had a grade of 0, 2, 3, 4
and 6 only. The total exposure of our observation is 111 ks for
each of the XIS detectors. We extracted the spectra and images
of the remnant from the 5 × 5 and 3 × 3 editing mode event
files using XSELECT version 2.4. For the spectral analysis we
generated the redistribution matrix file (RMF) and ancillary
response files (ARF) using XISRMFGEN and XISARFGEN
respectively. To analyze the spectral data we used the X-ray
9
It is also important to note that a fraction of XIS0 has not been functional
since 2009 June 23rd due to the damage caused by a micro meteorite. For more
information see: http://www.astro.isas.jaxa.jp/suzaku/doc/suzakumemo/
suzakumemo-2010–01.pdf.
10
http://www.astro.isas.jaxa.jp/suzaku/doc/suzakumemo/suzakumemo2007-08.pdf
11
http://heasarc.nasa.gov/docs/suzaku/processing/criteria xis.html
12
AtomDB 3.0.1 can be downloaded here: http://www.atomdb.org/
download.php
5
The Astrophysical Journal, 810:43 (12pp), 2015 September 1
Auchettl et al.
Figure 4. Suzaku XIS0 image of MSH 11–61A in the 0.5–7.0 KeV energy band. The image has been smoothed with a Gaussian function of width 0.′ 1 and has a
logarithmic scaling applied to it. Left: the green contours correspond to the 843 MHz MOST radio continuum emission used in Figure 1. Right: overlaid in green are
the regions we use for spectral extraction, described in Section 3.1.
Table 1
Best Fit Parameters for Our Background Model
Component
Parameter
Value
Cosmic X-ray background
NH (×1022 cm−2)
G
+1.07
1.470.83
1.40 (frozen)
Galactic ridge emission
NH (×1022 cm−2)
kT (keV)
Abundances
+0.08
0.610.06
0.22 ± 0.02
0.20 (frozen)
Galactic halo
NH (×1022 cm−2)
kT; (keV)
Ne
Mg
Si
+0.13
1.640.10
0.58 ± 0.02
+0.60
2.710.40
+0.23
2.310.18
+0.34
3.480.32
Reduced χ2 (dof)
1.80 (1602)
Figure 5. Sensitivity of our fits to changes in kTinit based on Dχ2. Here the red
plots correspond to kTinit = 5 keV, and the blue corresponds to kTinit = 20 keV.
The solid, dotted and dotted–dashed lines correspond to τ = 1010 s cm−3,
1011 s cm−3 and 1012 s cm−3 respectively. We assume kT = 0.5 keV and
nH = 1022 cm−2 for all spectra. From this plot we derive an upper limit for
kTinit of 2–5 keV corresponding to a shock velocity of 1300–2100 km s−1, for
which we then use in our fits.
Note. All uncertainties correspond to 1σ errors.
collisional equilibrium with an initial temperature kTinit that
suddenly drops to its final temperature kT. For our analysis, the
column density, ionization timescale, normalization, and final
temperature were left as free parameters. Due to the strong
emission lines from Mg, Si, and S the abundances of these
elements were also left free. Additionally we also let Ne and Fe
be free parameters for regions 2–3, as we found that varying
these significantly improved the fit. All other elemental
abundances were frozen to the solar values reported by Wilms
et al. (2000). The foreground absorbing column density NH was
modeled using TBABS (Wilms et al. 2000). Figure 7 shows an
example of the X-ray spectrum of MSH 11–61A as extracted
from region 3. The spectra derived for each region shown in
Figure 4 right panel all have similar features to the spectrum
shown in Figure 7.
We found that for regions 1–8 the fit favored an initial
temperature larger than the final temperature, while regions 9
and 10 favored an initial temperature smaller than the final
temperature. When left free, these initial temperatures would
hit the upper (or lower) limits of this parameter and the
associated abundances we obtained for our fits were unrealistically high (abundances of [Mg], [Si] > 10 relative to Wilms
et al. 2000). The ionization timescale for all regions was τ ∼
1012 s cm−3.
To investigate the sensitivity of our fits to values of kTinit,
we simulated spectra with similar counting statistics to those
from our regions of investigation, using kT = 0.5 keV,
NH = 1022 cm−2, and kTinit = 5 or 20 keV. We considered
cases for log t = 10, 11, and 12. We fit each spectrum to a
TBABS*VVRNEI model and then investigated the effect of
freezing kTinit values over a range from 0.1 to 100 keV. The
results are illustrated in Figure 5 where we plot Dχ2 versus
kTinit for spectra with actual kTinit values of 5 keV (red) and
20 keV (blue). Here the solid, dotted, and dashed lines
correspond, respectively, to log t = 10, 11, and 12. For low
6
The Astrophysical Journal, 810:43 (12pp), 2015 September 1
Auchettl et al.
Table 2
Spectral Fits for All 10 Individual Regions
Region
1
2
3
4
5
6
7
8
9
10
NH (1022 cm−2)
kT (keV)
kTinit (keV)
Ne
Mg
Si
S
Fe
τ (1012 s cm−3)
Reduced χ2
+0.12
1.760.15
+0.12
1.200.11
+0.14
1.320.13
+0.24
1.810.09
+0.10
1.870.17
+0.04
1.660.08
+0.21
1.670.16
+0.23
2.270.21
+0.45
2.000.30
+0.38
2.460.37
+0.03
0.430.02
+0.02
0.420.02
+0.03
0.400.02
+0.02
0.340.09
+0.06
0.270.01
+0.03
0.300.02
+0.05
0.360.03
+0.04
0.420.04
+0.22
0.800.27
+0.24
0.730.14
5
5
5
5
5
5
5
5
0.0808
0.0808
K
+0.12
0.230.09
+0.17
0.350.12
K
K
K
K
K
K
K
+0.52
3.230.46
+0.26
1.820.20
+0.33
2.040.26
+0.33
2.93-0.28
+0.33
2.490.28
+0.51
2.670.20
+1.07
4.110.77
+0.58
2.590.46
+2.77
3.631.02
+1.67
4.691.03
+1.07
6.320.92
+0.38
4.960.34
+0.50
4.650.42
+0.73
5.690.51
+0.74
5.21-0.75
+0.67
3.690.41
+1.92
7.051.30
+1.11
4.870.86
+5.72
6.91-1.96
+2.29
6.671.43
+0.24
1.920.19
+0.38
2.510.34
+0.43
2.620.36
+0.41
3.720.71
+1.52
3.341.14
+0.46
2.120.44
+1.42
4.271.11
+0.83
2.910.66
+5.20
3.061.17
+1.43
2.860.90
K
+0.08
0.110.06
+0.11
0.140.08
K
K
K
K
K
K
K
+0.19
1.590.14
+0.08
1.290.06
+0.06
1.260.06
+0.04
1.340.07
+0.09
1.280.11
+0.09
1.000.08
+0.08
1.090.07
+0.05
0.900.06
+0.22
0.22-0.11
+0.39
0.190.11
0.79
0.94
1.02
1.08
1.06
1.10
1.14
0.92
0.84
0.94
Note. All uncertainties correspond to the 90% confidence level.
ionization timescales the resulting fits are quite sensitive to the
kTinit value, while for timescales similar to that found in MSH
11–61A (τ ≈ 1012 s cm−3), the fits are insensitive to ktinit (i.e.,
the Dχ2 versus kTinit plot plateaus) for temperatures greater
than 2–5 keV. This is because above ∼2–5 keV, it is only the
emission from Fe and Ni that is significantly impacted by
higher kTinit values, because all other abundant ions are fully
stripped at these temperatures. Our observations do not have
enough counts around the Fe and Ni lines to provide sensitivity
to such an effect (though longer observations, particularly with
higher resolution, would provide such sensitivity; see below).
Based on where Dχ2 versus kTinit plot plateaus for τ ≈
12
10 s cm−3, we fix value of kTinit at 5 keV for all regions other
than 9 and 10. For the latter regions, whose fits indicate
kTinit < kT), we fix the initial temperature at the minimum
available value for the VVRNEI model (80.8 eV). In Table 2,
we list the best fit parameters for each of our spectra. All
uncertainties correspond to the 90% confidence level.
A recombining plasma that has an initial temperature of
2–5 keV implies that the shock must have had a velocity of
∼1300–2100 km s−1, assuming electron-ion equilibrium. This
high initial velocity suggests that the recombining plasma could
be the result of the SNR shock front initially expanding into a
dense CSM, reaching a high ionization state corresponding to
the high initial expansion velocity (and, thus, shock temperature). Such a scenario might result from expansion into an r−2
density profile characteristic of a stellar wind, with subsequent
expansion into the lower density regions resulting in rapid
cooling, leaving an overionized plasma (Itoh & Masai 1989;
Moriya 2012).
Even though our current data are unable to differentiate
between a plasma that has an initial temperature of 2–5 keV or
one that has initial temperature greater than this, with the
launch of ASTRO-H we will be able to differentiate between
these two cases. In Figure 6 we have plotted a simulated
background-subtracted spectrum that we would obtain in a 20
ks observation using the calorimeter on ASTRO-H for a plasma
that has an initial temperature of 2 keV and one that has an
initial temperature of 50 keV, assuming our fit parameters for
region 8. The background spectrum for ASTRO-H was obtained
from SIMX simulations.13 Here the black spectrum corresponds to a plasma that has an initial temperature of 2 keV (or
an initial velocity of ∼1300 km s−1), while the red spectrum
corresponds to a plasma that has an initial temperature of
13
Figure 6. A 20 ks background subtracted spectrum simulated for ASTRO-H for
a recombining plasma that has an initial temperature of 2 keV (black) and one
that has an initial temperature of 50 keV (red). The model was based on that
obtained for region 8 but with the initial temperature set to 2 and 50 keV. With
this observation one can easily differentiate between a recombining plasma that
has two different initial temperatures.
50 keV (or an initial velocity of ∼7000 km s−1, typical of the
high initial expansion speed of an SNR). One can see that with
a 20 ks ASTRO-H observation, we could easily differentiate
between two plasmas that have different initial temperatures.
Our spectra from MSH 11–61A are best described by a
recombining plasma model with the exception of emission
from regions 9 and 10, in the eastern and southeastern outskirts,
where an ionizing plasma is observed. Our results for the
central regions (1–3) agree well with those of Kamitsukasa
et al. (2015), who also find a recombining plasma for their
southwestern region, in agreement with our results. In contrast,
for their southeastern and northeastern regions, Kamitsukasa
et al. (2015) obtain best fits for an ionizing plasma. Given that
these regions combine emission from the eastern and southeastern protrusions (our regions 9 and 10), for which we also
observe an ionizing plasma, with emission from the outer
symmetric portion of the SNR, where we observe a
https://hea-www.harvard.edu/simx/
7
The Astrophysical Journal, 810:43 (12pp), 2015 September 1
Auchettl et al.
Table 3
Number Density of the X-Ray Emitting Material Estimated from the Best Fits
of the 10 Different Regions Shown in Figure 4
recombining plasma, the combination of two components may
explain the partial discrepancy. Conversely, Kamitsukasa et al.
(2015) report an ionizing plasma for their southwestern region,
which covers regions for which we obtain a recombining
plasma. This may result from our annular regions averaging
over multiple components.
All regions have an ionization timescale of
∼1.0 × 1012 cm−3 s, indicating that the X-ray emitting plasma
across the whole remnant is close to ionization equilibrium
(Smith & Hughes 2010). The temperature of the recombining
components ranges from 0.27 to 0.43 keV near the bright
central X-ray emission. Interestingly, the regions that are best
described by an ionizing plasma have the highest temperatures
out of all the regions that we analyzed, with temperatures of
0.80 and 0.73 keV respectively. On average, our derived
temperatures are lower than that derived by Kamitsukasa et al.
(2015), García et al. (2012) in their XMM–Newton analysis, and
Slane et al. (2002) in their ASCA analysis.
We find in all regions strong emission lines coming from
Mg, Si and S and all regions require super-solar abundances of
these elements. The enhancement of elemental abundances is
observed in many MM SNRs and indicates that some of the
X-ray emission we are observing arises from ejecta that have
been dispersed throughout the remnant and been mixed with
the swept up shocked material. Similar to Kamitsukasa et al.
(2015), and García et al. (2012) we also find an underabundance of Ne and Fe in regions 2–3. Unlike, Kamitsukasa
et al. (2015) we do not find evidence for overabundance of Ar
or the underabundance of O suggested by García et al. (2012).
When we freed these parameters we found that they do not
significantly improve our fit, thus we kept them at solar. Our
estimated abundances are slightly higher than that derived by
Kamitsukasa et al. (2015), García et al. (2012) and Slane et al.
(2002). This discrepancy arises from the fact that in our
analysis we use the abundance table by Wilms et al. (2000) and
the newly updated ATOMDB 3.0.1, while Kamitsukasa et al.
(2015) and García et al. (2012) use the table derived by Anders
& Grevesse (1989) and ATOMDB 3.0 and ATOMDB 2.0.2
respectively.
Our derived column density toward MSH 11–61A ranges
−2
+0.12
+0.38
22
between (1.200.11 - 2.46-0.37 ) ´ 10 cm . The column
density is highest in regions 8–10 which directly interacts with
the surrounding environment. Our estimates for NH are higher
than the column density derived by Slane et al. (2002), García
et al. (2012) and Kamitsukasa et al. (2015). This discrepancy
arises from the fact that we use a different abundance table and
a different absorption model.
Region
1
2
3
4
5
6
7
8
9
10
n
(d 7-0.5 f−0.5 cm−3)
+0.28
0.660.23
+0.24
0.710.22
+0.31
0.850.30
+1.17
1.28-0.43
+0.96
2.431.37
+1.01
2.681.20
+1.28
2.411.15
+2.82
5.872.32
+1.32
1.77-0.93
+1.41
2.231.17
of MSH 11–61A. Our density estimates for the bulk of the
remnant are consistent with the densities derived by Slane et al.
(2002) who attempted to reproduce the observed temperature
and brightness profiles of the remnant using the cloudy ISM
model by White & Long (1991) and a hydrodynamical model
that traces the evolution of the remnant, while incorporating the
effects of thermal conduction.
4. THE ORIGIN OF THE THERMAL X-RAY EMISSION
The total X-ray emitting mass in MSH 11–61A is given by
MX = 1.4nH mH fV , where mH is the mass of the hydrogen
atom, V is the volume from which the emission is observed,
and f is the filling factor. Using the estimated volumes and
derived densities for regions 1–7, we sum the masses to obtain
MX ≈ 480 d5/2f1/2 Me. This is comparable to the swept up
mass derived by Slane et al. (2002).
The enhancement of Mg, Si and S abundances throughout
the remnant suggests that the observed X-ray emission arises
in part from supernova ejecta. Assuming that all ejecta
have been shocked, we can estimate the mass of the
ejecta components based upon the measured abundances:
Mi = [(ai - 1) 1.4](ni nH )(mi mH ) MX where Mi is the ejecta
mass of species i, ai is its abundance relative to ISM
abundances, as listed in Table 2, mi is the atomic mass, and
ni/nH is its ISM abundance relative to hydrogen. We find that,
using the average of the measured abundances, the total ejecta
masses of Mg, Si, and S are, respectively, 0.37d 5 2f 1 2 M,
0.80d 5 2f 1 2 M, and 0.26d 5 2f 1 2 M. However, we note that
the abundances for Ne and Fe are both lower than ISM values,
meaning that we have no evidence for ejecta components for
these species, and suggesting caution in interpreting all of the
abundances. Taken at face value, however, the Mg, Si, and S
ejecta mass estimates are consistent with a progenitor mass
>25 Me (Thielemann et al. 1996).
Recombining plasma can arise from two main scenarios:
thermal conduction which is the rapid cooling of electrons due
to the interaction of the hot ejecta with the cold, dense
surrounding environment (Cox et al. 1999); or adiabatic
expansion which can occur when the SNR shockwave expands
through a dense circumstellar medium into a low density ISM
(Itoh & Masai 1989).
For thermal conduction to be the more likely scenario, the
recombining plasma is expected to be coincident with the
location of the MC, there should be a temperature decrease
toward the MC and one would expect the thermal conduction
3.2. Deriving the Density of the X-Ray Emitting Gas
The density of the X-ray emitting gas was calculated from
the normalization of the VVRNEI models using ne = 1.2nH.
We estimate the volume for each region by taking an area
equivalent to the extracted SNR regions shown in Figure 4
right panel and projecting this area through a filled sphere. The
estimated density n ≈ 1.1nH (assuming ISM abundances) is
listed for each region in Table 3.
+0.28
+2.82
The inferred density ranges from n=(0.660.23 - 5.87-2.32 )
−3
-0.5 −0.5
d7 f
cm and is highest in region 8 which is coincident
with the location of the dense MC found toward the west. The
density is lowest toward the center of the remnant where the
brightest X-ray emission is located, while the eastern part of the
remnant has a density that is intermediate of the central regions
8
The Astrophysical Journal, 810:43 (12pp), 2015 September 1
Auchettl et al.
observed γ-rays, which is plausible. In an attempt to
disentangle the likely source of the γ-ray emission, we have
plotted as the magenta contours in Figure 1 right panel, the HI
contours of the MC associated with MSH 11–61A. If one
assumes that the pulsar of ICG J11014-6103 can produce
significant γ-ray emission, we would expect the detection
significance peak to be skewed toward the position of ICG
J11014-6103 instead of in the direction of the remnant and MC
as is observed. Thus even though we cannot rule out that ICG
J11014-6103 is contributing to the γ-ray emission in Figure 3,
the association of the MC and the detection significance in this
region suggests that the emission most likely arises from the
interaction of the SNR with the MC, rather than ICG
J11014-6103.
timescale tcond to be less than or comparable to the age of the
remnant. We find recombining plasma in regions that are
directly interacting with the MC (regions 7 and 8), and we do
see a slight temperature decrease toward the MC based on the
annulus regions. The thermal conduction timescale is given
by (Spitzer 1962; Zhou et al. 2014) tcond ~ kne lT2 k
~ 56 (n e 1 cm-3)(lT 10 pc)2 (kT 0.6 keV)-5 2 (ln L 32) kyr,
where ne is the electron density and is calculated from our
best-fits listed in Table 2, lT is the scale length of the
temperature gradient, k is Boltzmann’s constant, κ is
the thermal conductivity for a hydrogen plasma and ln L is
the Coulomb logarithm. Assuming a distance of 7 kpc to the
remnant and a radius of ∼5′. 77, we calculate the length of
the temperature gradient to be ∼3.6 × 1019 cm. Using the
temperature and density of region 7 (see Tables 7 and 3
respectively), the thermal conduction timescale is estimated to
be ∼334 kyr. This is ∼16 times greater than the age of the
remnant (∼10–20 kyr), making it unlikely that the overionized
plasma arises via thermal conduction.
Another possibility is that the recombining plasma arises
from adiabatic cooling. To calculate trecomb, we use the best fit
ionization timescale for region 1–8 listed in Table 2 and divide
these by the electron density of each region. We obtain a
recombining timescale between ∼(2–125) d 7-0.5f−0.5 kyr, which
is comparable to the age of MSH 11–61A, making this scenario
the most likely. This is consistent with the results reported by
Kamitsukasa et al. (2015) and the velocity implied by our upper
limit for kTinit.
5.2. Modeling the Broadband Emission of MSH 11–61A
To investigate the nature of the broadband emission from
MSH 11–61A we use a model that calculates the non-thermal
emission produced by a distribution of electrons and protons.
The π0 decay model is based on the proton–proton interaction
model by Kamae et al. (2006), with a scaling factor of 1.85 for
helium and heavy nuclei as suggested by Mori (2009). The
leptonic emission models are based on those presented by
Baring et al. (1999) and Bykov et al. (2000) for the
synchrotron/IC and non-thermal bremsstrahlung emission
mechanisms. We assume a spectral distribution of our
accelerated particles to be
⎛ p ⎞
dNi
= ai p-ai exp ⎜ ⎟,
dp
⎝ p0 i ⎠
5. THE NATURE OF γ-RAY EMISSION
5.1. Pulsar Contribution and the
Integral Source ICG J11014-6103
(1 )
where i is the particle species, ai is the normalization of the
particle distribution, αi is the particle distribution index which
is equal to (1-G)/2, where G is the photon index and p0i is the
exponential particle momentum cut-off. This distribution is
transformed from momentum space to energy space such that
the exponential cut-off is defined by an energy input, E0i. The
sum of the integrals of each spectral distribution is set to equal
the total energy in accelerated particles within the SNR shell,
ECR = òESNR, where ò is the efficiency of the SNR in
depositing energy into CRs.
Pulsars that are detected within the Fermi-LAT energy band
(see the second Fermi-LAT Pulsar catalog by Abdo
et al. 2013), have spectra that is well characterized by a power
law with an exponential cut-off of 1–5 GeV. As the γ-ray
spectrum of MSH 11–61A can be described using an
exponential cut-off of Ecut ∼ 4.2 GeV, we still need to consider
the scenario that the γ-ray emission we observe arises from a
nearby pulsar other than J1105-6107.
Using the Australian Telescope National Facility (ATNF)
Pulsar Catalogue (Manchester et al. 2005), there are 9 pulsars
including J1105-6107 within 5°, whose spin down power is
sufficient to produce the γ-ray flux of MSH 11–61A. All of
these pulsars, except for J1105-6107, are >1° from the centroid
of the γ-ray emission making it unlikely that any of these
pulsars are contributing significantly to the observed emission
of MSH 11–61A. As we removed the contribution of J11056107 from the γ-ray data as described in Section 2.1, we can
also rule out its contribution.
Recently, Pavan et al. (2014) investigated the nature of the
X-ray and radio emission of the INTEGRAL source ICG
J11014-6103, which they call the lighthouse nebula. In X-rays
this nebula exhibits a prominent jet-like feature that is
perpendicular to an elongated cometary tail, and a point
source. The source of this X-ray structure is a neutron star
travelling supersonically and we have plotted its position as the
white X shown in Figure 1. This neutron star has a spin down
power of E˙ ~ 10 37 erg s−1. As the γ-ray luminosity of MSH
11–61A is 2.5 × 1035 erg s−1, the PSR of ICG J11014-6103
would require an efficiency of L g E˙ ~ 2.5% to produce the
5.2.1. Hadronic Origin of the Observed γ-Rays
In Figure 8, we have plotted the model fits to the broadband
emission of MSH 11–61A. The radio spectrum is a combination of multiple observations by Milne et al. (1989), Whiteoak
& Green (1996), and Filipovic et al. (2005). The X-ray upper
limit was derived by fitting a power law with a photon index
similar to that of RX J1713.7-3946 (Uchiyama et al. 2003),
Kepler and RCW 86 (Bamba et al. 2005) (G = 2.3), to our
models of the Suzaku data. The upper limit corresponds to the
flux in which the additional non-thermal component begins to
affect our reduced χ2. The solid magenta line corresponds to
the π0-decay model that adequately reproduces the observed γray spectrum of MSH 11–61A. We have also plotted as the
purple dotted–dashed line the synchrotron model that sufficiently reproduces the radio spectrum, assuming an electron–
proton ratio (Kep) of 0.01, while the IC model falls below the
plotted axis. For completeness we have also plotted the nonthermal bremsstrahlung contribution as the orange dashed line.
Table 4 lists the parameters which reproduce the π0-decay,
9
The Astrophysical Journal, 810:43 (12pp), 2015 September 1
Auchettl et al.
Figure 7. Suzaku XIS0, XIS1 and XIS3 spectrum extracted from region 3 in Figure 4. These spectra were fitted simultaneously using our background model described
in Table 1 and an absorbed VVRNEI model as described by the parameters listed in Table 2. The spectra are overlaid with the best fit model with their chi2 residuals.
The spectra derived for each region shown in Figure 4 right panel all have similar features as the spectrum shown here.
Figure 8. Broadband fits to the non-thermal radio emission (green data points), non-thermal X-ray upper limit derived from the Suzaku data (dark green limit) and the
Fermi-LAT γ-ray emission (as described in Figure 3) of MSH 11–61A. The pion decay, non-thermal bremsstrahlung and synchrotron models defined by the
parameters in Table 4 are shown as the solid magenta line, dashed orange line and dotted–dashed purple line, respectively. The corresponding IC model falls below the
plotted axis.
Table 4
Model Parameters and Density Estimates for the Pion Decay and Leptonic Model for MSH 11-61A
Object
MSH 11-61A
Distance
(kpc)
αproton
αelec
E0proton
(GeV)
E0elec
(GeV)
Magnetic Field
(μG)
Ambient Density
n0 (cm−3)
X-Ray Density
n (cm−3)
7.00
4.39
3.15
6.05
6.05
28
9.20
see Table 3
particles to escape the emission volume (Malkov
et al. 2011, 2012).
As non-thermal X-ray emission has not been observed from
MSH 11–61A, we are unable to constrain the cut off energy of
the electron population. Thus to model the radio emission of
MSH 11–61A we assume the electron distribution has the same
cut-off energy as the proton distribution. We are able to
reproduce the radio spectrum using an electron distribution that
has a power law index of αe = 3.15 and a magnetic field of
28 μG. The magnetic field implied by the synchrotron
synchrotron emission, IC and non-thermal bremsstrahlung
models plotted in Figure 8.
A π0-decay model arising from a proton distribution with a
power law index of αp = 4.34 and a cut-off energy of
E0p ~ 6.05 GeV, adequately described the γ-ray spectrum of
MSH 11–61A. The cut-off energy of the proton spectrum
derived in this model is much smaller than the TeV cut-off one
would expect for protons (Reynolds 2008). This could indicate
that due to the high density of the surrounding environment,
efficient CR acceleration is suppressed allowing accelerated
10
The Astrophysical Journal, 810:43 (12pp), 2015 September 1
Auchettl et al.
modeling is larger than the magnetic field of the ISM
(∼3–5 μG). This enhancement could arise from magnetic field
amplification due to the compression of the surrounding
medium by the SNR shock-wave.
In the π0-decay model, the calculated γ-ray flux is
proportional to the ambient density of the surrounding ambient
medium and the total proton energy. Assuming a conservative
upper limit of 40% of the total supernova explosion energy
goes into accelerating CRs, we can estimate the density of the
γ-ray emitting material. For our π0-decay model of MSH
11–61A we obtain an ambient density of 9.15 cm−3. Similar to
many other SNRs that exhibit hadronic emission (e.g., W41,
MSH 17-39, G337.7-0.1: Castro et al. 2013; Kes 79: Auchettl
et al. 2014), this density is much larger than the ambient
density estimate derived from our X-ray analysis (see Table 3).
This discrepancy could arise from the SNR shock-wave
interacting with dense clumps of material that are cold enough
such that do not radiate significantly in X-rays (Castro &
Slane 2010; Inoue et al. 2012). If these clumps have a high
filling factor, then the densities that we derive in our X-ray
analysis would underestimate the mean local density. Our
inferred ambient densities as well as the association of MSH
11–61A with a MC toward the west of the remnant supports the
conclusion that MSH 11–61A is interacting with dense material
that does not radiate in X-rays.
An alternative scenario is that the enhanced γ-ray emission
arises from highly energetic particles escaping the acceleration
region and are interacting with dense gas upstream of the shock
(e.g., Aharonian & Atoyan 1996; Gabici et al. 2008; Lee et al.
2008 and Fujita et al. 2009). However, a majority of these
particles come from the high-energy portion of the γ-ray
spectrum and the observation of “low” energy γ-rays may lead
to inconsistencies with this scenario.
is interacting with a MC toward the north and southwest.
Similar to previous studies, the inferred density from our pion
decay model is much higher than that implied by the thermal
X-ray emission. Suzaku data reveal that the bulk of the X-ray
emission of MSH 11–61A arises from a single recombining
plasma with enhanced abundances of Mg, Si and S with some
regions also requiring an underabundance of Ne and Fe, while
the emission toward the east of the remnant arises from an
ionizing plasma with Mg, Si and S. The origin of the
recombining plasma is most like adiabatic cooling. We find
that the results from our central regions (1–3) and our regions
9–10, agree well with those that Kamitsukasa et al. (2015)
obtained for their corresponding regions. The enhancement of
Mg, Si and S suggests that some of the observed emission
arises from shocked ejecta and that the progenitor of MSH
11–61A had a mass >25 Me.
REFERENCES
Abdo, A. A., Ackermann, M., Ajello, M., et al. 2010a, ApJ, 722, 1303
Abdo, A. A., Ackermann, M., Ajello, M., et al. 2010b, Sci, 327, 1103
Abdo, A. A., Ajello, M., Allafort, A., et al. 2013, ApJS, 208, 17
Ackermann, M., Ajello, M., Albert, A., et al. 2012, ApJS, 203, 4
Ackermann, M., Ajello, M., Allafort, A., et al. 2013, Sci, 339, 807
Aharonian, F. A., & Atoyan, A. M. 1996, A&A, 309, 917
Anders, E., & Grevesse, N. 1989, GeCoA, 53, 197
Auchettl, K., Slane, P., & Castro, D. 2014, ApJ, 783, 32
Bamba, A., Yamazaki, R., Yoshida, T., et al. 2005, ApJ, 621, 793
Baring, M. G., Ellison, D. C., Reynolds, S. P., et al. 1999, ApJ, 513, 311
Brand, J., & Blitz, L. 1993, A&A, 275, 67
Bykov, A. M., Chevalier, R. A., Ellison, D. C., et al. 2000, ApJ, 538, 203
Caraveo, P. A. 1993, ApJL, 415, L111
Castro, D., & Slane, P. 2010, ApJ, 717, 372
Castro, D., Slane, P., Carlton, A., et al. 2013, ApJ, 774, 36
Cordes, J. M., & Lazio, T. J. W. 2002, arXiv:astro-ph/0207156
Cox, D. P., Shelton, R. L., Maciejewski, W., et al. 1999, ApJ, 524, 179
Dickel, J. R. 1973, ApL, 15, 61
Dwarkadas, V. V. 2005, ApJ, 630, 892
Ellison, D. C., Patnaude, D. J., Slane, P., et al. 2010, ApJ, 712, 287
Ergin, T., Sezer, A., Saha, L., et al. 2014, ApJ, 790, 65
Filipovic, M. D., Payne, J. L., & Jones, P. A. 2005, SerAJ, 170, 47
Foster, A. R., Ji, L., Smith, R. K., et al. 2012, ApJ, 756, 128
Fujita, Y., Ohira, Y., Tanaka, S. J., et al. 2009, ApJL, 707, L179
Gabici, S., Casanova, S., & Aharonian, F. A. 2008, in AIP Conf. Proc. 1085,
High Energy Gamma-Ray Astronomy, ed. F. A. Aharonian, W. Hofmann &
F. Rieger (Melville, NY: AIP), 265
Gaisser, T. K., Protheroe, R. J., & Stanev, T. 1998, ApJ, 492, 219
García, F., Combi, J. A., Albacete-Colombo, J. F., et al. 2012, A&A, 546, A91
Goss, W. M., Radhakrishnan, V., Brooks, J. W., et al. 1972, ApJS, 24, 123
Halpern, J. P., Tomsick, J. A., Gotthelf, E. V., et al. 2014, ApJL, 795, L27
Henley, D. B., & Shelton, R. L. 2013, ApJ, 773, 92
Hobbs, G., Lorimer, D. R., Lyne, A. G., et al. 2005, MNRAS, 360, 974
Inoue, T., Yamazaki, R., Inutsuka, S., et al. 2012, ApJ, 744, 71
Itoh, H., & Masai, K. 1989, MNRAS, 236, 885
Kamae, T., Karlsson, N., Mizuno, T., et al. 2006, ApJ, 647, 692
Kamitsukasa, F., Koyama, K., Uchida, H., et al. 2015, PASJ, 67, 16
Kaneda, H., Makishima, K., Yamauchi, S., et al. 1997, ApJ, 491, 638
Kaspi, V. M., Bailes, M., Manchester, R. N., et al. 1997, ApJ, 485, 820
Kaspi, V. M., Manchester, R. N., Johnston, S., et al. 1996, AJ, 111, 2028
Kawasaki, M., Ozaki, M., Nagase, F., et al. 2005, ApJ, 631, 935
Kawasaki, M. T., Ozaki, M., Nagase, F., et al. 2002, ApJ, 572, 897
Kesteven, M. J., & Caswell, J. L. 1987, A&A, 183, 118
Kesteven, M. J. L. 1968, AuJPh, 21, 369
Koyama, K., Tsunemi, H., Dotani, T., et al. 2007, PASJ, 59, 23
Kramer, M., Bell, J. F., Manchester, R. N., et al. 2003, MNRAS, 342, 1299
Kushino, A., Ishisaki, Y., Morita, U., et al. 2002, PASJ, 54, 327
Lee, S.-H., Kamae, T., & Ellison, D. C. 2008, ApJ, 686, 325
Malkov, M. A., Diamond, P. H., & Sagdeev, R. Z. 2011, NatCo, 2, 194
Malkov, M. A., Diamond, P. H., & Sagdeev, R. Z. 2012, PhPl, 19, 082901
Manchester, R. N., Hobbs, G. B., Teoh, A., et al. 2005, AJ, 129, 1993
McClure-Griffiths, N. M., Dickey, J. M., Gaensler, B. M., et al. 2005, ApJS,
158, 178
5.2.2. Leptonic Emission of the Observed γ-Rays
For IC scattering to be the dominant mechanism producing
the γ-rays of MSH 11–61A, we would require greater than the
entire kinetic explosion energy just in electrons, assuming that
the electron to proton ratio is similar to that measured at Earth
(Kep ∼ 0.01) and that this emission arises from a non-thermal
population of electrons being accelerated by the shock-front.
This makes it difficult to conclude that IC scattering is the
dominant mechanism producing the observed γ-rays.
For non-thermal bremsstrahlung to dominate the GeV
emission we require a Kep > 0.2, assuming the maximum
density derived from our X-ray analysis (see Table 3). Local
measurements imply Kep ∼ 0.01 (Gaisser et al. 1998), while
γ-ray modeling of other SNRs imply Kep < 0.01 (e.g., Ellison
et al. 2010), making it unlikely that non-thermal bremsstrahlung emission is the dominant emission mechanism.
6. CONCLUSION
Seventy months of Fermi-LAT data reveal significant (∼5σ)
γ-ray emission from SNR MSH 11–61A. This emission is
consistent with being located along or inside the western limb
of the remnant given the angular resolution of the Fermi-LAT
and is adjacent to regions that show a strong recombining
plasma component. By modeling the broadband spectrum, we
find that the emission is best described by a hadronic model,
while a leptonic scenario is energetically unfavorable. This is
consistent with CO and H I observations that indicate the SNR
11
The Astrophysical Journal, 810:43 (12pp), 2015 September 1
Auchettl et al.
Mills, B. Y., Slee, O. B., & Hill, E. R. 1961, AuJPh, 14, 497
Milne, D. K., Caswell, J. L., Kesteven, M. J., et al. 1989, PASAu, 8, 187
Mori, M. 2009, APh, 31, 341
Moriya, T. J. 2012, ApJL, 750, L13
Nolan, P. L., Abdo, A. A., Ackermann, M., et al. 2012, ApJS, 199, 31
Pavan, L., Bordas, P., Pühlhofer, G., et al. 2014, A&A, 562, A122
Reynolds, S. P. 2008, ARA&A, 46, 89
Reynoso, E. M., Johnston, S., Green, A. J., et al. 2006, MNRAS, 369, 416
Rho, J., & Petre, R. 1998, ApJL, 503, L167
Rosado, M., Ambrocio-Cruz, P., Le Coarer, E., et al. 1996, A&A, 315, 243
Sato, T., Koyama, K., Takahashi, T., et al. 2014, PASJ, 66, 124
Seward, F. D. 1990, ApJS, 73, 781
Slane, P., Bykov, A., Ellison, D. C., et al. 2015, SSRv, 188, 187
Slane, P., Smith, R. K., Hughes, J. P., et al. 2002, ApJ, 564, 284
Smith, R. K., Brickhouse, N. S., Liedahl, D. A., et al. 2001, ApJL, 556, L91
Smith, R. K., & Hughes, J. P. 2010, ApJ, 718, 583
Spitzer, L. 1962, Physics of Fully Ionized Gases (2nd ed.; New York:
Interscience)
Tawa, N., Hayashida, K., Nagai, M., et al. 2008, PASJ, 60, 11
Thielemann, F.-K., Nomoto, K., & Hashimoto, M.-A. 1996, ApJ, 460, 408
Uchiyama, Y., Aharonian, F. A., & Takahashi, T. 2003, A&A, 400, 567
Uchiyama, Y., Aharonian, F. A., Tanaka, T., et al. 2007, Natur, 449, 576
Warren, J. S., Hughes, J. P., Badenes, C., et al. 2005, ApJ, 634, 376
White, R. L., & Long, K. S. 1991, ApJ, 373, 543
Whiteoak, J. B. Z., & Green, A. J. 1996, A&AS, 118, 329
Wilms, J., Allen, A., & McCray, R. 2000, ApJ, 542, 914
Yu, M., Manchester, R. N., Hobbs, G., et al. 2013, MNRAS, 429, 688
Zhou, P., Safi-Harb, S., Chen, Y., et al. 2014, ApJ, 791, 87
12